Observation of Enhanced X-Ray Emission from the CTTS AA Tauri During One Transit of an Accretion Funnel ﬂow

A&A 475, 607–617 (2007) Astronomy DOI: 10.1051/0004-6361:20078084 & c ESO 2007 Astrophysics

Observation of enhanced X-ray emission from the CTTS AA Tauri during one transit of an accretion funnel ﬂow

N. Grosso1, J. Bouvier2, T. Montmerle2, M. Fernández3, K. Grankin4, and M. R. Zapatero Osorio5

1 Observatoire astronomique de Strasbourg, Université Louis-Pasteur, CNRS, INSU, 11 rue de l’Université, 67000 Strasbourg, France e-mail: [email protected] 2 Laboratoire d’astrophysique de Grenoble, Université Joseph-Fourier, CNRS, INSU, 414 rue de la Piscine, 38041 Grenoble, France 3 Instituto de Astrofísica de Andalucía, CSIC, Camino Bajo de Huétor 50, 18008 Granada, Spain 4 Ulugh Beg Astronomical Institute of the Uzbek Academy of Sciences, Astronomicheskaya 33, 700052 Tashkent, Uzbekistan 5 Instituto de Astrofísica de Canarias (IAC), vía Láctea s/n, 38205 La Laguna, Tenerife, Spain Received 14 June 2007 / Accepted 27 August 2007

ABSTRACT

Context. Classical T Tauri stars are young solar-type stars accreting material from their circumstellar disks. Thanks to a favorable inclination of the system, the classical T Tauri star AA Tau exhibits periodic optical eclipses as the warped inner disk edge occults the stellar photosphere. Aims. We intend to observe the X-ray and UV emission of AA Tau during the optical eclipses with the aim of localizing these emitting regions on the star. Methods. AA Tau was observed for about 5 h per XMM-Newton orbit (2 days) over 8 successive orbits, which covers two optical eclipse periods (8.22 days). The XMM-Newton optical/UV monitor simultaneously provided UV photometry (UVW2 filter at 206 nm) with a ∼15 min sampling rate. Some V-band photometry was also obtained from the ground during this period in order to determine the dates of the eclipses. Results. Two X-ray and UV measurements were secured close to the center of the eclipse (∆V ∼ 1.5 mag). The UV flux is the highest just before the eclipse starts and the lowest towards the end of it. UV flux variations amount to a few 0.1 mag on a timescale of a few hours and up to 1 mag on a timescale of a week, none of which are correlated with the X-ray flux. We model it with a weekly modulation (inner disk eclipse), plus a daily modulation, which suggests a non-steady accretion, but needs a longer observation to be confirmed. No such eclipses are detected in X-rays. Within each 5 h-long observation, AA Tau has a nearly constant X-ray count rate. On a timescale of days to weeks, the X-ray flux varies by a factor of 2−8, except for one measurement where the X-ray count rate was nearly 50 times higher than the minimum observed level even though photoelectric absorption was the highest at this phase, and the plasma temperature reached 60 MK, i.e. a factor of 2−3 higher than in the other observations. This X-ray event, observed close to the center of the optical eclipse, is interpreted as an X-ray flare. Conclusions. We explain the variable column density with the low-density accretion funnel flows blanketing the magnetosphere. The lack of X-ray eclipses indicates that X-ray emitting regions are located at high latitudes. Furthermore, the occurrence of a strong X-ray flare near the center of the optical eclipse suggests that the magnetically active areas are closely associated with the base of the high-density accretion funnel flow. We speculate that the impact of this free-falling accretion flow onto the strong magnetic field of the stellar corona may boost the X-ray emission. Key words. stars: individual: AA Tauri – stars: pre-main sequence – X-rays: stars – accretion, accretion disks

1. Introduction because the X-ray luminosity is proportional to the plasma emission measure, which scales linearly with the plasma volume but T Tauri stars (TTSs), i.e. young (1–10 Myr) solar-type stars, also with the square of the plasma electronic density. That most are conspicuous X-ray emitters. Their high X-ray luminosities 28−31 −1 27 −1 of TTSs’ X-ray emission arises in an active magnetic corona (LX 10 erg s ) compared to the Sun (LX 10 erg s is supported by direct Zeeman measurements on photospheric at solar maximum) and intense flaring activity (up to LX 32−33 −1 spectral lines which indicate surface magnetic fields of a few 10 erg s ) make them appear as extremely active young kilogauss (e.g., Guenther et al. 1999; Johns-Krull et al. 1999, suns in the X-ray domain. The analogy with solar activity has 2004; Johns-Krull 2007; Yang et al. 2005). However, the dy- been quite successful in ascribing the X-ray emission of TTSs namo mechanism producing the magnetic field in these fully to an optically thin, magnetically confined coronal plasma in − convective stars is still being discussed (e.g., Preibisch et al. collisional equilibrium at temperatures of 10 100 MK, emit- 2005; Briggs et al. 2007). ting a thermal bremsstrahlung continuum and emission lines (see, e.g., review by Feigelson & Montmerle 1999). As ob- The solar paradigm, where the X-ray emitting plasma is con- served in the X-ray coronae of active stars (e.g., Ness et al. fined in a magnetic loop with both feet anchored in the stellar 2004), the enhanced X-ray luminosity of TTSs can easily be photosphere, has been questioned in the context of a Classical explained by coronal structures with high plasma density, TTS (CTTS), since it accretes material from its circumstellar disk. Inside a few stellar radii above the CTTS surface, the stel- Figures 2, 3, and 9, and Appendix A are only available in electronic lar magnetic field pressure is higher than the ram pressure of the form via http://www.aanda.org accreting gas. As a result, the stellar magnetosphere truncates the

Article published by EDP Sciences and available at http://www.aanda.org or http://dx.doi.org/10.1051/0004-6361:20078084 608 N. Grosso et al.: Enhanced X-ray emission from AA Tauri inner accretion disk and controls the accretion flows. The gas is Table 1. Journal of the XMM-Newton observations of AA Tau (PI: mainly accreted from the disk edge to the stellar surface along J. Bouvier). the dominant large-scale stellar magnetic lines, creating accretion funnel flows. The free-falling gas hits the stellar surface at Obs. Rev. ObsId Feb. 2003a Exposurea the feet of the accretion funnel flows, where the kinetic energy (d) (h) is dissipated in a shock, producing a hot excess emission (see 1 583 015680201 14.10–14.29 4.7 review on magnetospheric accretion by Bouvier et al. 2007b). 2 584 015680301 16.13–16.33 4.7 3 585 015680501 18.07–18.26 4.7 The X-ray grating spectrometers aboard Chandra and 4 586 015680401 20.02–20.22 4.7 XMM-Newton are able to obtain spectra of the X-ray brightest 5 587 015680601 22.13–22.35 5.4 CTTSs, where the emission line triplets of He-like elements are 6 588 015680701 24.03–24.25 5.3 resolved. This provides a powerful tool for assessing the elec- 7 589 015680801 26.52–26.71 4.6 tronic density of the X-ray emitting plasma (e.g., Porquet et al. 8 590 015680901 28.30–28.50 4.7 2001). In several CTTSs, plasma with high electronic density a 13 −2 The observation beginning, end, and duration is given for MOS1. (ne ∼ 10 cm ) and low temperature (∼3 MK), atypical of stellar coronae, were identified and therefore attributed to accretion shocks (Kastner et al. 2002; Stelzer & Schmitt 2004; distribution to reconcile the high value of the X-ray absorption Schmitt et al. 2005; Robrade & Schmitt 2006; Argiroffi et al. outside the eclipse and the low optical extinction. 2007; Günther et al. 2007; Telleschi et al. 2007), whereas some In Sect. 2, we present the X-ray and UV properties of CTTSs display no such evidence (Audard et al. 2005; Smith et al. AA Tau based on reanalysis of the full data set provided by 2005; Güdel et al. 2007b). these XMM-Newton observations. In particular, we report X-ray During the Chandra Orion Ultradeep Project (COUP, see (pn+MOS1+MOS2) and UV light curves with a time resolution Getman et al. 2005b), where the TTSs of the Orion nebula clus- of ∼15 min. We show that a bright and hot flare was observed ter were monitored almost continuously over ∼13 days with the during the second observation. In Sect. 3, thanks to our ground- Advanced CCD Imaging Spectrometer (ACIS), numerous X-ray based observations and optical OM data, we determine the dates flares were observed. In a few cases, a size of several stellar radii of the optical eclipses. This allow us to compare the UV and was derived for the magnetic loop confining the X-ray emitting X-ray variabilities versus the rotational phase. We show that the plasma, large enough to connect the stellar surface with the edge UV minimum is outside the eclipse and that the (confirmed) vari- of the accretion disk (Favata et al. 2005). ation in the column density are not correlated with the rotational We propose to directly constrain the source and location of phase. In Sect. 4, we propose another origin for this variable col- the X-ray emission in CTTSs by using eclipses. Eclipse map- umn density, and we discuss the origin of the X-ray flare that was ping has been successfully used in binary active stars to recon- observed during an optical eclipse. struct the coronae (e.g., Güdel et al. 2001, 2003) or to localize the flaring plasma (e.g., Schmitt & Favata 1999; Schmitt et al. 2. X-ray and UV properties of AA Tau 2003). Our target star is AA Tau, located in the Taurus molecular cloud complex at a distance of ∼140 pc (e.g., Kenyon et al. We present the XMM-Newton observations and event selections 1994). AA Tau is quite a typical member of the CTTS class, in Sect. 2.1. We report the X-ray light curves in Sect. 2.2, and with a K7 spectral type, a bolometric luminosity of ∼0.8 L, the X-ray spectra and plasma parameters in Sect. 2.3. We com- a stellar mass of ∼0.8 M, and a stellar radius of ∼1.85 R. pare these results with previous X-ray observations of AA Tau It exhibits moderate accretion disk diagnostics (near-IR excess, in Sect. 2.4. The OM observations are presented in Sect. 2.5: optical veiling, Balmer line emission), but with the remarkable the UV light curves are reported in Sect. 2.5.1 and modeled in property of being viewed nearly edge-on. Bouvier et al. (1999, Sect. 2.5.2; and the optical photometry, obtained with the field 2003, 2007a) have reported evidence of a modulation of the acquisition exposure, is reported in Sect. 2.5.3. photospheric flux and spectroscopic diagnostics with a period of 8.22 days, corresponding to the rotational period of the star. This stellar flux modulation has been interpreted as the peri- 2.1. XMM-Newton observations and event selections odic eclipse of the stellar photosphere by the optically thick, Our observational strategy with XMM-Newton was to cover two magnetically-warped inner disk edge located at 8.8 stellar radii. consecutive modulation periods (∼17 days) in order to demon- Photopolarimetric variations confirm the presence of an opti- strate the reproducibility of the phenomenon from one rotational cally thick wall located at the disk edge and periodically eclips- cycle to the next. The temporal sampling of the X-ray light curve ing the stellar photosphere (Ménard et al. 2003). does not need to be very dense because AA Tau spends nearly We obtained 8 observations of AA Tau with XMM-Newton the same amount of time being eclipsed as being entirely vis- (Jansen et al. 2001), which allows simultaneous observations ible. Therefore, we requested one 4 h-exposure with EPIC pn with an EPIC pn (Strüder et al. 2001) and two EPIC MOS per XMM-Newton orbit (2 days) over 8 successive orbits. We (Turner et al. 2001) X-ray spectroimaging cameras, and the used the full-frame science mode of the EPIC cameras with Optical/UV monitor (OM; Mason et al. 2001). We supplemented the medium optical blocking filter. The pointing nominal co- these X-rays and UV observations with an optical ground-based ordinates were 04h34m55s.5, 24◦2854. 0 (J2000 equinox). The monitoring of AA Tau to secure the dates of the optical eclipses. journal of the XMM-Newton observations of AA Tau is given in These XMM-Newton observations were previously reported in Table 1. Schmitt & Robrade (2007, hereafter SR), who used the mini- The data reduction was made using the XMM-Newton mum of the UV light curve as a proxy of the optical eclipse. SR Science Analysing System (SAS, version 7.0.0) For each obser- find variable X-ray absorption, “such that the times of maximal vation, the event lists for each camera of EPIC were produced X-ray absorption and UV extinction coincide”, and introduce using the SAS tasks epchain and emchain. We used the back- an additional absorption in a disk wind or a peculiar dust grain ground lightcurves computed by these tasks in the 7.0−15.0 keV N. Grosso et al.: Enhanced X-ray emission from AA Tauri 609

Fig. 1. XMM-Newton observations of AA Tau. Top and bottom panels show the background-subtracted UV (180–250 nm) and X-ray (0.5–7.3 keV) light curves obtained with the Optical/UV Monitor (OM) and EPIC (pn+MOS1+MOS2), respectively. In each panel, the label indicates the time interval used to bin the light curve. Grey hourglasses and crosses data points indicate UV photometry obtained with the large and small OM central imaging window. The X-ray count rates are corrected from (circular) aperture. The gray stripes indicate the beginning and the end of the EPIC observation. energy range to determine the time intervals affected by back- circular aperture photometry using the fraction of PSF counts ground proton flares. More than half of the observing time inside the (circular) extraction region calculated by the SAS, as- was affected by bad weather in space. For each instrument, we suming a fixed photon energy of 1.5 keV.Finally, the light curves made from the low background time intervals a sky image with of the three detectors were summed to produce the EPIC light 3-pixels in the 0.5−7.3 keV energy range1. The X-ray counter- curves. We estimated the missing pn data at the beginning of the part of AA Tau was detected in all the observations. observations by multiplying the MOS1+MOS2 count rates by For the X-ray light curves of AA Tau, we selected the 1.17, the median scaling factor between MOS1+MOS and pn in source+background events within a circular region centered on the second observation where AA Tau was the brightest. AA Tau. The extraction position and radius were optimized to The bottom panel of Fig. 1 shows the EPIC light curves of maximize the signal-to-noise ratio in the sky image. The back- AATauinthe0.5−7.3 keV energy range. Within each 5 h-long ground in MOS and pn was extracted using an annular region observation, AA Tau exhibited a nearly constant X-ray flux. The centered on AA Tau and a box region at the same distance to minimum X-ray flux, which we call hereafter the quiescent level, the CCD readout node, respectively, where areas illuminated by was observed during the observation #5, where the averaged other weak X-ray sources were excluded. For the X-ray spectra, − EPIC count rate was 0.030 ± 0.002 counts s 1. On a timescale we used the same extraction regions and only those time inter- of days to weeks, the X-ray flux varies by a factor of 2−8, ex- vals with low background; we selected events with energy above cept between the end of the first observation and the beginning 0.3 keV and the usual stronger selection criteria2. We computed of the second observation, where the EPIC count rate jumped in the corresponding redistribution matrix files and ancillary re- − less than 2 days from 0.093 ± 0.003 to 1.42 ± 0.01 counts s 1, sponse files. i.e. a level 47 ± 3 times stronger than the quiescent level. Then, the EPIC count rate decayed in less than 2 days to 0.075 ± − 2.2. X-ray light curves 0.003 counts s 1, i.e. a level only 2.5 ± 0.2 times stronger than the quiescent level. For each instrument and observation, we first built the source+background and the background light curves with 1 s Such large amplitudes in the X-ray fluxes of young stellar time bins starting at the first good time interval (GTI) of MOS1. objects are usually observed during X-ray flares, which have We rebinned the light curve to 900 s to increase the signal. Then, typical light curves with fast rise and peak phase and slower we subtracted the background light curve scaled to the same (exponential) decay phase, associated with fast heating and slow source extraction area from the source+background light curve. cooling of the magnetically confined plasma (e.g., Imanishi et al. From the GTI extension, we used an IDL routine to scale 2003; Favata et al. 2005). X-ray flares with unusually long rise up count rates and errors affected by any lost of observing time, phases have also been reported (e.g., Grosso et al. 2004; Wang mainly due to triggering the counting mode during high-flaring et al. 2007; Broos et al. 2007). background periods, when the count rate exceeded the detec- For comparison, we note that, when assuming a typical con- tor telemetry limit. We also corrected count rates and errors for vertion ratio between XMM-Newton/EPIC and Chandra/ACIS-I count rates (e.g., Ozawa et al. 2005), the average EPIC 1 We selected single, double, triple, and quadruple pixel events (i.e. count rate of AA Tau, ∼0.08 counts s−1 at the distance of PATTERN in the 0 to 12 range) and also applied the predefined filter ∼ #XMMEA_EM for MOS. the Taurus molecular cloud (d 140 pc), would convert to ∼ / −1 2 For pn, we selected only single and double pixel events (i.e. 0.001 Chandra ACIS-I counts s at the distance of the Orion PATTERN in the 0 to 4 range) with FLAG value equal to zero; for MOS, nebula (d ∼ 450 pc). AA Tau, put in the Orion nebula cluster, we selected PATTERN in the 0 to 12 range and applied the predefined would have then been brighter than 67% of the X-ray sources filter #XMMEA_SM. in the COUP. To check whether the behavior of the light curve 610 N. Grosso et al.: Enhanced X-ray emission from AA Tauri of AA Tau is consistent with the one of an X-ray flare, we can compare it to the light curve data set obtained in the COUP. We then use the COUP results of the Bayesian block (BB) variability analysis (developed by Scargle 1998; adapted for the COUP data set and coded in IDL by one of us, N.G.), which seg- mented the X-ray light curves into contiguous sequences of constant count rates (see Getman et al. 2005b). We define a source to be variable, if there is more than one BB; with BBmin and BBmax the minimum and the maximum count rate levels, respectively (see, e.g., Stassun et al. 2007, for an application of COUP time-averaged X-ray variability). The latter and the former are viewed as the quiescent level and the peak level of the brightest flare, respectively. When this criterion is applied, there are 977 variable sources (out of 1616 COUP sources). We find only 20 variable sources with peak amplitude (BBmax/BBmin) and duration larger than 45 and 4.7 h, respectively, as observed for Fig. 4. EPIC pn (black), MOS1 (red), MOS2 (green) spectra of AA Tau for observation #2. The lines show our best fit using two temperature AA Tau. For comparison, our subsample sources have BBmax = −4− −1 plasma combined with photoelectric absorption (Table 2). The residuals 10 2 counts s ; i.e., they are at peak between 200 times are plotted in sigma units with error bars of size one. fainter and 100 times brighter than AA Tau (at the distance of the Orion nebula) at peak. Then, we make a visual examination of the selected BB light d.o.f.) to 617.2 (for 597 d.o.f.). Given the new and old values of curves to eliminate sources with peak flare BB having a spuri- χ2 and number of degrees of freedom, an F-test indicates a prob- ous long duration (including for example a passage through the ∼ −8 / ability of 10 for the null hypothesis; therefore, we conclude van Allen belts) and or sources with a decay phase too slow to that it is reasonable to add this extra temperature component to reproduce the level observed at the beginning of the third obser- improve the fit. Our best fit with a two-temperature plasma is vation of AA Tau. For an exponential decay phase with a decay τ τ < . showninFig.4. timescale d, this latter criteria is equivalent to d 0 6day. Table 2 gives the corresponding plasma parameters. On Finally, we exhibit the brightest flares from COUP 313, 874, and timescale of 2 days, the photoelectric absorption of the X-ray 970, which fullfill our criteria (see online Fig. 2). These sources spectra is not constant, but the observed relative variations are are bona fine members of the Orion nebula cluster (Getman et al. lower than a factor of two. The minimum and maximum values 2005a). of the corresponding column density are ∼1.0 × 1022 cm−2 and We conclude that a bright X-ray flare with a rapid cooling 1.8 × 1022 cm−2, respectively; the latter was observed during the phase can reproduce the large amplitude well, along with the second observation4. SR did not specify the model of X-ray ab- flatness of the light curve of AA Tau at its maximum. The fol- sorption that they used, but their one-temperature plasma model lowing X-ray spectra analysis confirms this interpretation. gave qualitatively similar results. We stress that the value of the column density can even be increased by 50% when revised so- 2.3. X-ray spectra and plasma parameters lar abundances (i.e., metal poor) are adopted for the abundances of the absorbing material (see online Appendix A). For each observation the pn, MOS1, and MOS2 spectra were The plasma temperature is ∼23 MK during the low activ- binned to 25 counts per spectral bin. X-rays are detected up to ity levels. The observations showing an increase in X-ray count 10 keV (see online Fig. 3). The spectra are featureless, except in rates (namely #2, #6, and #7) also correspond to phases with the second observation where a prominent line around 6.7 keV the highest plasma temperatures (60 MK, 40 MK, and 45 MK, is visible both in the pn and MOS spectra, corresponding to the respectively) and therefore to flaring activity. Fe XXV triplet emission line (Fig. 4). During the second observation, the X-ray surface flux, i.e., The pn, MOS1, and MOS2 spectra were fitted simultane- the intrinsic X-ray luminosity divided by the stellar surface, ously with XSPEC (version 12.3.0; Dorman & Arnaud 2001) to peaked to 108 erg s−1 cm−2. The temperatures of the hot (60 MK) derive the plasma parameters. The model is an X-ray emission and cool (15 MK) plasma components, associated with this ele- spectrum from collisionally-ionized diffuse gas, output from vated level of X-ray surface flux, are consistent with the ones the Astrophysical Plasma Emission Code (vapec3), which in- observed in the most active TTSs of COUP (see Fig. 11 of cludes continuum and emission lines. The plasma element abun- Preibisch et al. 2005). The cool plasma component is usually dances were fixed to typical values measured in the coronae observed in the coronae of active stars and may define a fun- of young stars with grating X-ray spectroscopy (see online damental coronal structure, which is probably related to a class appendix Table A.1), to allow a direct comparison with the of compact loops with high plasma density (see Preibisch et al. XMM-Newton Extended Survey of the Taurus molecular cloud 2005, and references therein). (XEST; Güdel et al. 2007a). This emission model is combined We make an estimate of the hot loop length using the method with the wabs photoelectric absorption model, which is based on of Reale et al. (1997) (see also Favata et al. 2005, for an appli- the photoionization cross sections of Morrison & McCammon cation to COUP data). Assuming a peak temperature of 60 MK (1983) and the solar abundances of Anders & Ebihara (1982). and a flare-decay time lower than 0.6 day (see Sect. 2.2), we find The online Fig. 3 shows our best fits. For the second observation, for the semi-circular loop a half-length lower than 7.1 R for a we find a better fit by adding a second temperature component, which reduced the χ2 from 655.4 (for 599 degrees of freedom; 4 The increase of the spectrum slope between 1 and 2 keV observed between observations #1 and #2, is not only due to the increase of the 3 More information can be found at: absorption, as argued in SR, but also for 20% to the increase in the http://hea-www.harvard.edu/APEC/sources_apec.html plasma temperature. N. Grosso et al.: Enhanced X-ray emission from AA Tauri 611

Table 2. Best parameters of simultaneous ﬁtting of EPIC pn, MOS1, and MOS2 spectra with XSPEC.

Plasma temperature Emission measure Goodness-of-ﬁt Fluxb Intrinsic luminosity ηc

a 2 Obs. Rate NH T1 T2 EM1 EM2 χν (ν) Q 0.5–2, 2–8, 0.5–8 keV (cnts s−1)(1022 cm−2)(MK)(1053 cm−3)(%)(1030 erg s−1) 1 0.057 0.95 ...... 25.9 ...... 1.0 0.93(50) 61 1.9 0.6 0.4 0.9 –3.5 ±0.002 0.87–1.02 ...... 23.9–28.4 ...... 0.9–1.0 2 0.800 1.54 ...... 43.9 ...... 14.0 1.09(599) 5 37.8 8.2 9.2 17.4 –2.2 ±0.008 1.51–1.57 ...... 42.7–45.2 ...... 13.6–14.3 2 0.800 1.80 14.6 59.4 9.8 9.3 1.03 (597) 28 37.9 11.8 9.5 21.3 –2.2 ±0.008 1.75–1.85 13.4–15.7 54.6–66.0 8.8–10.7 7.8–10.8 3 0.046 1.36 ...... 20.8 ...... 1.1 1.21(30) 20 1.6 0.7 0.3 1.0 –3.5 ±0.002 1.23–1.51 ...... 18.4–23.4 ...... 1.0–1.3 4 0.031 1.33 ...... 21.5 ...... 0.8 0.79(25) 76 1.1 0.5 0.2 0.7 –3.6 ±0.002 1.17–1.52 ...... 18.7–24.8 ...... 0.6–0.9 5 0.019 1.04 ...... 22.6 ...... 0.4 1.32(15) 18 0.7 0.2 0.1 0.4 –3.9 ±0.002 0.85–1.27 ...... 18.3–28.8 ...... 0.3–0.5 6 0.050 1.54 ...... 40.5 ...... 0.9 1.20(39) 18 2.3 0.5 0.5 1.1 –3.5 ±0.002 1.40–1.72 ...... 34.6–47.4 ...... 0.8–1.0 7 0.147 1.02 ...... 44.8 ...... 2.1 1.05(110) 33 6.5 1.2 1.4 2.6 –3.1 ±0.005 0.97–1.08 ...... 41.3–48.6 ...... 2.0–2.2 8 0.043 1.53 ...... 23.3 ...... 1.2 1.19(29) 23 1.8 0.7 0.4 1.1 –3.4 ±0.002 1.39–1.69 ...... 20.7–26.5 ...... 1.0–1.4 Notes: We fit the X-ray spectra with the continuum and emission lines produced by an optically thin plasma in thermal collisional ionization equilibrium model (vapec). We use, for the plasma element abundances, typical values observed in the coronae of young stars with fine X-ray spectroscopy (Güdel et al. 2007a, see Table A.1). The wabs photoelectric absorption model uses the photoionization cross sections of Morrison & McCammon (1983) and the solar abundances of Anders & Ebihara (1982). Errors are given at the 68% confidence level (i.e. ∆χ2 = 1.0for each parameter of interest), which corresponds to 1σ for Gaussian statistics. For observation #2, a better fit is obtained using a plasma with two temperature components. a pn count rate (0.2–12 keV). b Observed X-ray flux (0.5–8.0 keV) in unit of 10−13 erg cm−2 s−1. c Logarithm of the X-ray intrinsic luminosity (0.5–8 keV) to the bolometric luminosity (0.8 L) ratio. freely decaying loop with no heating, or lower than 1.6 R for Walter & Kuhi (1981) reported a detection with 0.030 ± a strongly sustained heating5 in the cooling phase. Therefore, 0.005 IPC counts s−1, during the first observation with Einstein the height of the semi-circular loop is lower than 1−4.5 R. (bright enough to exhibit a crude spectrum). But, nearly one Consequently, this loop cannot connect the stellar surface and year later with a 5 times longer observation, Walter & Kuhi the inner accretion disk, distant by 7.8 stellar radii, and is likely (1984) reported only an upper limit of 0.004 IPC counts s−1. anchored in the stellar photosphere. Walter & Kuhi (1984) associated the X-ray emission detected The high temperature and X-ray surface flux observed dur- from AA Tau as a quiescent level and interpreted the non- ing the second observation point to an enhanced X-ray activity detection in the framework of the smothered coronae of TTSs produced by a bright X-ray flare. Such bright flares on the Sun (see Walter & Kuhi 1981), where it was proposed that mass are sometimes associated with coronal mass ejection (CME). ejection could increase the absorbing column density enough to Therefore, we cannot rule out that the maximum of column den- smother the coronal (quiescent) X-ray emission. Neuhäuser et al. sity observed during the second observation is due to this ener- (1995) reported a RASS detection 10 years later with 0.014 ± getic event. 0.006 PSPC counts s−16. AA Tau was also detected during the PSPC pointed observation with 0.022 ± 0.005 PSPC counts s−1 (see in the WGA catalogue of ROSAT point sources, the source 2.4. Comparison with previous X-ray observations of AA Tau 1WGA J0434.9+2428, located 37 off-axis; White et al. 1996), AA Tau was previously observed several times in X-rays at dif- but (again) only during the shortest exposure. ferent epochs: on March 4, 1980 and February 7, 1981 with A further constraint can be derived on the X-ray variabil- the IPC aboard Einstein, with an exposure of 0.6 h and 2.8 h, ity of AA Tau by combining these multiple-epoch observations respectively; on August 10, 1990 during the ROSAT All Sky with our better assessment of its quiescent level thanks to bet- Survey (RASS) with the PSPC aboard ROSAT, with an exposure ter spectra. For a plasma temperature ∼23 MK and an unab- of 0.2 h; and on February 22 and August 18, 1993 during two sorbed X-ray luminosity in the energy band from 0.5 to 8.0 keV pointed PSPC observations of the young binary Haro 6-13 (PI: of ∼0.5 × 1030 erg s−1 (equivalent to an unabsorbed X-ray H. Zinnecker), with an exposure of 0.6 h and 1.5 h, respectively. flux of 2.1 × 10−13 erg s−1 cm−2), absorbed by a column density of ∼1022 cm−2, we computed with PIMMS7 that Einstein/IPC 5 ζ Reale et al. (1997) use , the slope of the flare decay in a log-log (0.2−4.5 keV) and ROSAT/PSPC (0.12−2.48 keV) would both diagram of the plasma temperature vs. the squared-root of the emission observe only ∼0.002 counts s−1. This low count rate is consistent measure (a proxy of the plasma density), as the diagnostic to assess the level of sustained heating in the analysis of stellar flares. For the XMM-Newton energy coverage, the corresponding range for ζ values 6 Note that SR reported a RASS rate that is 10 times higher than the are 0.4 and 1.9, for strongly sustained heating and freely decaying loops, one reported by Neuhäuser et al. (1995). respectively. 7 Available at http://asc.harvard.edu/toolkit/pimms.jsp 612 N. Grosso et al.: Enhanced X-ray emission from AA Tauri with the IPC upper limit8, and we found that it is twice lower than the background level (inside a 1.5-radius circle) at the location of AA Tau in the second PSPC pointed observation. Therefore, the previous non-detections in X-rays are consistent with the quiescent level observed with XMM-Newton. We conclude that the previous X-ray detections with Einstein, the RASS, and ROSAT/PSPC, were made during high levels of activity, where AA Tau was about 18, 8, and 13 times, respectively, above its quiescent level.

2.5. XMM-Newton optical/UV monitor light curves 2.5.1. UV photometry The OM was operated in the imaging mode default, which uses 5 imaging windows plus a small (1.7 × 1.7) central imaging window. The imaging mode default consists of a sequence of 5 exposures where one of the 5 imaging windows covers a large fraction of the OM field-of-view (17 × 17) with 1 × 1 spatial resolution, while the small central imaging window ensures a continuous monitoring of the target at the center of the field- Fig. 5. Weekly and daily UV variability of AA Tau observed with the of-view with 0. 5 × 0. 5 spatial resolution. (For an illustration OM. Top panel: the dotted and dashed-dotted lines show the weekly of this exposure sequence and a light curve obtained with the modulation of the UV flux attributed to the eclipse period (8.22 days) and the overall modulation of the UV flux. The horizontal line shows the small central window see Fig. 85 of the XMM-Newton Users’ average flux. Bottom panel: the dashed line shows the daily modulation Handbook and Grosso et al. 2007, respectively.) We requested of the UV folded in phase after subtraction of the weekly modulation. the minimum available exposure time in imaging mode default Symbols are the same in both panels. The formula of the overall fit (800 s) to monitor any change in the UV photometry (UVW2 fil- (dashed-dotted line) is given by Eq. (1). ter at 206 nm) with a ∼15 min sampling rate. We ran the OM imaging mode pipeline. We made our own mainly due to fast variations on a time scale of a day, which can- IDL program to plot the light curves of AA Tau from the source not be properly reproduced with this simple model that only in- lists of the small central window (*OM*SWSRLI0000.FIT), and cludes weekly variations. A Lomb periodogram of the residuals the large central window (*OM*SWSRLI1000.FIT). The latter suggests the presence of an additional (high-frequency) modula- only × provided 1 data point per OM exposure sequence (about 5 tion with a period of 0.87 ± 0.03 day. 800 s), which was obtained simultaneously with the first (out of We then computed a grid of fits in the amplitude and phase 5) data point of the small central window. In a few exposures, parameter space to look for optimized values that maximize the AA Tau is missing in the observation source list of the small cen- peak of the periodogram residuals. This is equivalent to a simul- tral window, but a visual inspection of the corresponding images taneous fitting of both modulations with the long-term period confirms the detection. Therefore, we completed the light curve fixed. We find a better fit (χ2 = 103.8 for 77 d.o.f.). Given the by doing aperture photometry. new and old values of χ2 and number of degrees of freedom, a The top panel of Fig. 1 shows the UV light curves of AA Tau. F-test indicates a probability of ∼10−11 for the null hypothesis; We note that the UV light curves of AA Tau reported in SR were therefore, we conclude that it is reasonable to add this extra mod- limited to the UV photometry obtained with the large OM cen- ulation to improve the fit. Moreover, a Lomb periodogram of the tral imaging window (see for comparison the grey hourglasses residuals shows no other modulations and a perfect substraction in the top panel of our Fig. 1 and their Fig. 1). The continuous of the two modulations. The fit formula is given by the equation monitoring with the central window is crucial for accurately de- termining the UV variations on a one-hour timescale. We find UV = 16.30 ± 0.02 mag that AA Tau is variable in UV by an amount of a few 0.1 mag +(0.20 ± 0.02) cos{2π[t − (17.2 ± 0.1)]/8.22} on a timescale of a few hours, and up to 1 mag on a timescale +(0.24 ± 0.02) cos{2π[t − (14.75 ± 0.02)]/(0.87 ± 0.04)}, (1) of a week. There are no correlations between the UV and X-ray variations. where t is the day in February 2003, and the errors are given at the 68% confidence level. The top panel of Fig. 5 shows the UV light curve and the fit 2.5.2. Modeling of the UV variability given by Eq. (1). The bottom panel shows the UV light curve The UV excess observed in CTTSs is attributed to the accretion folded in phase with the high-frequency period after the sub- shocks at the base of the accretion funnel flows (see review on straction of the long-term modulation. This light curve looks magnetospheric accretion by Bouvier et al. 2007b). Therefore, fairly convincing. However, a longer UV observation with the we assume that the UV flux must be modulated somehow by the OM is necessary for a definitive confirmation of this high- warped inner disk, and we perform a least-square fitting of the frequency period, which would suggest a non-steady accretion. UV light curve using a cosine function with free amplitude and phase and with a period fixed to 8.22 days (Bouvier et al. 2007a). 2.5.3. Optical variability from the field acquisition exposure The resulting fit (χ2 = 206.5 for 80 d.o.f.) is not acceptable, We propose here to obtain extra information on the optical pho- 8 However, the RASS rate is not consistent with the IPC upper limit. tometry of the target thanks to the OM Field Acquisition expo- Note that SR argued the opposite. sures (FAQs), where several stars are also optically detected in N. Grosso et al.: Enhanced X-ray emission from AA Tauri 613

Table 3. Journal of ground-based optical observations.

Feb. 2003 Site Tel. Observer Nobs (d) 13.89–28.88 Teide (Spain) 0.8 m M.R. Z. 177 16.83–17.96 Sierra Nevada (Spain) 1.5 m M. F. 5 20.63–28.64 Mt Maidanak (Uzbek.) 0.5 m K. G. 8

Differential photometry was performed on CCD images and absolute photometry from photomultiplier observations, with an accuracy on the order of 0.01 mag. Somewhat larger systematic errors (≤0.05 mag) might result from the relative calibration of the photometry between sites. All data reduction procedures can be found in Bouvier et al. (2003). The top panel of Fig. 7 shows the ground-based optical light curve; for comparison, the other panels show the UV and X-ray light curves and the plasma parameters. Unfortunately, due to bad weather conditions in Europe, no ground telescope was able to obtain simultaneous optical observations with XMM-Newton. Fig. 6. Optical light curves of the OM guide stars obtained with the field In the most favorable cases, ground optical photometry was obtained only 4.3 h, 3.1 h, 2.8 h, and 2.0 h before the begin- acquisition exposure. The dashed lines and grey stripes indicate the raw − count rates of guide stars and one-sigma variations, respectively. The ning of the MOS observations #1 #4, respectively, and only thick line and error bars show the optical light curve of AA Tau. Labels 3.3 h after the end the last MOS observation. Despite the lim- indicate counterparts in the Guide Star Catalog (version 2.3.2). ited time sampling, a dip, very likely associated with a primary eclipse of AA Tau, was detected, however, at the beginning of the monitoring campaign. This supports what we the OM field of view and are then available as comparison stars deduced previously in Sect. 2.5.3 from the optical FAQ images for variability studies. that XMM-Newton observations #1 and #2 were obtained during The FAQ is a short exposure (10 s) V-filter image always the eclipse. Consequently, the FAQ images can also be safely taken at the start of each science observation (i.e., at the be- used to identify an eclipse during XMM-Newton observations #5 ginning of the MOS observation) to allow proper identification and #6, located within the Feb. 23−27, 2003 gap of the ground of guide stars and to compensate for small pointing errors. A observation. thresholding is applied to the FAQ image by the OM onboard software to identify guide stars. The derived offset is then ap- 3.2. Folded light curves and eclipse phases plied to the OM science windows. The information (positions in detector coordinates, count rates, etc.) on the guide stars is To better determine the dates of the optical eclipses, we com- reported in the *OMS40000RFX* data files delivered with the pared our photometry with those obtained just 6 months apart, Observation Data Files (ODFs). Instead of scheduling an ad- from Aug. 27 to Oct. 24, 2003 (Bouvier et al. 2007a). The top ditional OM exposure in the imaging mode default with the panel of Fig. 8 shows both photometry data sets folded together V-filter, the use of the FAQ image allowed us to observe twice as in phase using the 8.22-day rotational period (Bouvier et al. long with the UVW2 filter, by saving at least 5 × 800 s (without 2007a); i.e., we introduced no phase difference between the two taking time overheads into account). data sets. The (arbitrary) common origin of phase was chosen to In each observation, we recovered sky coordinates from de- match the light curve of season 2003 shown in Fig. 15 of Bouvier tector coordinates and identified the counterparts of guide stars et al. (2007a), taking the eclipse center at phase 0.5. in the Guide Star Catalog (version 2.3.2). Then, we built a light The behavior of our photometry is consistent with the shape curve for the 13 guide stars (including AA Tau). Figure 6 shows of the light curve observed during the season 2003. We can ex- the variations in the raw count rates in the optical of the guide clude any phase offset larger than 0.05 between the two epochs. stars during our campaign. All guide stars, except AA Tau, ex- A deep primary eclipse (∆V ∼ 1.5 mag) is well detected at the hibited small (lower than 2σ) relative variations in the optical. beginning of our monitoring campaign. There is also some ev- By contrast, AA Tau exhibited large relative variations, in par- idence of a shallow secondary eclipse detected only at the end ticular, during observations #1, #2, #5, and #6 where a large dip of our monitoring campaign after the last XMM-Newton obser- is visible. We conclude that these 4 XMM-Newton observations vation. The differences in brightness of about 0.5 and 1 mag were most likely made during the optical eclipses of AA Tau. between the Feb. and Aug. 2003 data, observed at the beginning and the end of the primary eclipse, respectively, can be explained with a longer duration of the eclipse in Feb. 2003. The 3. UV and X-ray variabilities versus rotational phase brightening on Feb. 16, 2003, close to the center of the primary ∼ 3.1. Ground-based optical photometry eclipse, was observed with a high sampling rate ( 2min)and exhibited only a fast decay. Therefore, this event is similar to the The journal of ground-based optical observations is given transient brightening events usually observed in the faint state in Table 3. Observations were carried out from three sites of AA Tau (Bouvier et al. 1999). The primary eclipses were cen- over two weeks, covering our XMM-Newton observations in tered on Feb. 15.5 and Feb. 23.5, 2003, and covered phases rang- February 2003, using either CCD detectors or a photomultiplier ing from ∼0.3 to ∼0.7. Our ground-based monitoring confirms tube (Mt Maidanak). Measurements were obtained in the V filter. that XMM-Newton observations #1, #2, #5, and #6 were made 614 N. Grosso et al.: Enhanced X-ray emission from AA Tauri

Fig. 7. Light curves and plasma parameters of AA Tau. The panels show from top to bottom: the optical, UV, and X-ray light curves (labels indicate observation numbers); the variations in the plasma emission measure(s) and temperature(s), and the photoelectric absorption (Table 2). The top horizontal axis indicates the corresponding phase for the rotation period of 8.22 days. The (arbitrary) phase origin is from Bouvier et al. (2007a). during primary eclipses. In particular, observations #2 and #6 exclude a secondary eclipse at the start of observations #4 and were secured close to the center of the primary eclipse. We can #8 thanks to the FAQ, and moreover the corresponding X-ray N. Grosso et al.: Enhanced X-ray emission from AA Tauri 615

Fig. 8. Light curves and plasma parameters of AA Tau folded in phase with the rotation period of 8.22 days. Symbols and the phase origin are identical to the ones used in Fig. 7. In the top panel, black dots show the optical ground monitoring from Aug. 27 to Oct. 24, 2003 (Bouvier et al. 2007a) for comparison. The horizontal arrow shows our estimate of the eclipse phase based on the optical ground monitoring. light curves show no decay. Therefore, we conclude that the was observed during the XMM-Newton observation #3 at the shallow secondary eclipse probably started after the end of our end of the egress phase, i.e., outside the primary eclipse. Our observations. model of the UV flux variations helps to disentangle the weekly The UV flux is at its highest when the primary eclipse starts modulation (produced by the warped disk) and daily varia- and lowest towards the end of it. Indeed, the lowest UV flux tions. Equation (1) indicates that the weekly modulation was at 616 N. Grosso et al.: Enhanced X-ray emission from AA Tauri minimum on 2003 Feb. 17.2 (see also the top panel of Fig. 5), band9, the plasma parameters can be derived from spectral fit- which corresponds to a phase delay of about 0.2 compared to ting, which provides, in particular, a better measurement of the the optical eclipse. Therefore, the warped disk produces a max- column density and the X-ray luminosity. imum of obscuration of the UV flux at the end of the optical We find that the column density, derived from the photo- eclipse. Consequently, the true maximum of the UV flux oc- electric absorption of the X-rays emitted by the active corona of 22 −2 22 −2 curred around phase 0.2, well before the start of the eclipse. We AA Tau, varies from NH ∼ 1.0 × 10 cm to 1.8 × 10 cm . did not survey this time interval with XMM-Newton, but we note However, the optical extinction of AA Tau by the dust is very that a brightening in the B-band of AA Tau was observed around low with AV = 0.78 mag (Bouvier et al. 1999), which can be 21 −2 phase 0.2 during the 1995 campaign (Bouvier et al. 1999). The converted to NH ∼ (1.2 ± 0.1) × 10 cm , using the rela- 21 −2 −1 delay between the optical and UV eclipses and the lower depth tion NH/AJ = (5.6 ± 0.4) × 10 cm mag (Vuong et al. of the UV eclipse (∆UVW2 ∼ 0.40 mag), compared to the opti- 2003), combined with the extinction law of Rieke & Lebofsky cal eclipse (∆V ∼ 1.5 mag), suggests a trailing accretion funnel (1985), AJ = 0.282 × AV (for their adopted RV value of 3.1). flow, producing a strong absorption of the UV photons emitted We conclude that there is an excess of column density, varying at the accretion shock. from 0.9 × 1022 to 1.7 × 1022 cm−2. These values are consis- No eclipses were detected in X-rays. The variable photo- tent with the one reported by SR. To explain this excess of col- electric absorption of the X-ray spectra is not correlated with umn density, SR introduced an additional absorption in a disk the rotational phase; indeed, similar low and high values of the wind (with no dust, and cool enough to avoid producing any column density were observed both during the eclipse and out- soft X-ray emission), or a peculiar dust grain distribution with ∼ . side it. The maximum of the column density was observed at RV 0 4 to reconcile the observed extinction and X-ray absorp- phase 0.6, close to the center of the optical eclipse, during the tion. However, we note that three-dimensional MHD simulations X-ray flare. However, increases in the column density were also of disk accretion to a rotating magnetized star (e.g., Romanova reported during large stellar flares, and solar flares are some- et al. 2003) show that the stellar magnetosphere is far from be- times associated with coronal mass ejection (see the review on ing an empty cavity. Matter accretes mainly through narrow fun- ∼ 12 −3 X-ray astronomy of stellar coronae by Güdel 2004, and refer- nel loaded with high-density material ( 10 cm ), which are ences therein). Therefore, we cannot rule out that the peak in col- surrounded by lower density funnel flows that blanket nearly the umn density is due to this energetic event. The gas column den- whole magnetosphere. Therefore, we propose explaining this ex- sity on the line of sight produced by the warped disk is around cess of gas by this lower density funnel flows filling the stellar 1025 cm−2 (Bouvier et al. 1999), which is high enough to absorb magnetosphere. Dividing the excess column density by the width all the X-rays emitted by AA Tau, even during a bright and hot of the magnetosphere of AA Tau (about 7.8 stellar radii) leads to 10 −3 flare. The lack of eclipses in X-rays indicates that X-ray emit- a density of a few 10 cm , which is compatible with this inter- ting regions are located above the high-density disk warp at high pretation. Moreover, the multiple spirals visible in the simulated latitudes. accretion flows should help to produce a variable column density. This low density gas, located below the radius of dust sub- The online Fig. 9 shows the scattered plot of the average OM limation (close to the inner accretion disk for AA Tau), is then count rate versus the column density, where the count rate error dust free. includes the observed variations in the UV flux (see Fig. 1). The / ffi − . Accreting TTSs show on average a LX Lbol ratio 2.5 times correlation coe cient of this sample is 0 72, suggesting a true lower than in non-accreting TTSs (see for COUP and XEST re- correlation between the two physical parameters as argued by sults, Preibisch et al. 2005; and Briggs et al. 2007, respectively). SR, who used smaller error bars for the count rates. However, we This is generally interpreted as either a direct or indirect con- note that the data obtained outside the eclipse are located within sequence of the accretion process (see discussion in Preibisch the cloud of points, with error bars covering all the range of ob- et al. 2005; Telleschi et al. 2007; Gregory et al. 2007). The served values for column density and count rate. Therefore, the median value of log(LX/Lbol)forAATauis−3.5; for compar- column-density variation cannot be attributed to the disk warp. ison this value is intermediate between −3.7and−3.3, the me- Assuming a frequency of one X-ray flare per 650 ks, as dian values found for accretors and non-accretors, respectively observed in young solar-mass stars (Wolk et al. 2005), our (Preibisch et al. 2005; Briggs et al. 2007). The peculiar orienta- 38.8 h-exposure should only detect 0.2 flare from AA Tau. The tion of AA Tau may help to partly overcome the extinction by center of the optical eclipse is limited to about 0.2 in phase. the accretion funnels (Gregory et al. 2007). Therefore, the combined probability of observing an X-ray flare Recently, Johns-Krull (2007) has reported new magnetic by chance during the center of the transit of an accretion fun- field measurements for CTTSs, based on Zeeman broadening of nel flow is only 0.04. Moreover, the median flare level in young photospheric absorption lines in the near-IR, showing that the solar-mass stars is only 3.5 times the characteristic level (Wolk observed mean magnetic field is in all cases greater than the et al. 2005), whereas we observed a flare with a larger ampli- field predicted by pressure equipartition arguments. Johns-Krull tude. This suggests that this event is associated with the magnetic (2007) suggests that the very strong fields on these stars decrease area corresponding to the base of the dipolar magnetic field line, the efficiency with which convective gas motions in the photo- which controls the accretion funnel flow. sphere can tangle magnetic loops in the corona. For AA Tau,

9 To convert the ROSAT count rate of AA Tau to X-ray luminosity, Neuhäuser et al. (1995) estimate an energy conversion factor from the visual extinction (converted to foreground hydrogen column density) 4. Discussion and conclusions and the second hardness ratio (sensitive to the plasma temperature). This leads to an X-ray luminosity of 0.4 × 1030 erg s−1 (see Table 1 The high throughput of XMM-Newton allows us to obtain spec- of Johns-Krull 2007), which is likely to be underestimated by a factor tra of AA Tau at each phase of its activity. In contrast to previous of ten (see our simulations of count rates in Sect. 2.4), because the col- ROSAT observations, where the measurement for this source was umn density of AA Tau (see Table 2, and discussion below) is about limited to count rate and hardness ratios in the soft X-ray energy 10 times higher than the one assumed by Neuhäuser et al. (1995). N. Grosso et al.: Enhanced X-ray emission from AA Tauri 617

Johns-Krull (2007) reports a value of 2.78 kG for the mean mag- Dorman, B., & Arnaud, K. A. 2001, in Astronomical Data Analysis Software netic field, which is 2.7 times higher than the field strength at and Systems X, ed. F. R. Harnden, Jr., F. A. Primini, & H. E. Payne, ASP pressure equipartition; and predicts from solar scaling an (unob- Conf. Ser., 238, 415 × 30 −1 Favata, F., Flaccomio, E., Reale, F., et al. 2005, ApJS, 160, 469 served) X-ray luminosity of about 4 10 erg s .(Forcom- Feigelson, E. D., & Montmerle, T. 1999, ARA&A, 37, 363 parison this value is 10 times greater than the minimum level Getman, K. V., Feigelson, E. D., Grosso, N., et al. 2005a, ApJS, 160, 353 that we observed.) However, our observation shows that a bright Getman, K. V., Flaccomio, E., Broos, P. S., et al. 2005b, ApJS, 160, 319 X-ray flare can occur during the transit of an accretion funnel Gregory, S. G., Wood, K., & Jardine, M. 2007, ArXiv e-prints, 704 Grosso, N., Montmerle, T., Feigelson, E. D., & Forbes, T. G. 2004, A&A, 419, flow and locate the active X-ray area probably close to the foot of 653 the accretion funnel flow. We speculate that a magnetic interac- Grosso, N., Audard, M., Bouvier, J., Briggs, K. R., & Güdel, M. 2007, A&A, tion exists between the free-falling accretion flow and the strong 468, 557 magnetic field of the stellar corona, which may trigger magnetic Güdel, M. 2004, A&ARv, 12, 71 reconnections and give rise to bright X-ray flares, which in turn Güdel, M., Audard, M., Magee, H., et al. 2001, A&A, 365, L344 Güdel, M., Arzner, K., Audard, M., & Mewe, R. 2003, A&A, 403, 155 boost the X-ray emission. Güdel, M., Briggs, K. R., Arzner, K., et al. 2007a, A&A, 468, 353 This campaign of observations of AA Tau with XMM- Güdel, M., Skinner, S. L., Mel’Nikov, S. Y., et al. 2007b, A&A, 468, 529 Newton shows that X-ray spectroscopy with CCD provides a Guenther, E. W., Lehmann, H., Emerson, J. P., & Staude, J. 1999, A&A, 341, unique tool for probing the circumstellar dust-free gas in this ob- 768 Günther, H. M., Schmitt, J. H. M. M., Robrade, J., & Liefke, C. 2007, A&A, ject and that some magnetic flares are likely to be associated with 466, 1111 the accretion process. Longer coordinated optical and X-ray/UV Imanishi, K., Nakajima, H., Tsujimoto, M., Koyama, K., & Tsuboi, Y. 2003, observations of AA Tau are still needed to obtain continuous PASJ, 55, 653 light curves, which is crucial for confirming the flaring behavior Jansen, F., Lumb, D., Altieri, B., et al. 2001, A&A, 365, L1 of active regions associated with the accretion funnel flow, and Johns-Krull, C. M. 2007, ApJ, 664, 975 Johns-Krull, C. M., Valenti, J. A., & Koresko, C. 1999, ApJ, 516, 900 daily modulation of the UV flux. Simultaneous Zeeman-Doppler Johns-Krull, C. M., Valenti, J. A., & Saar, S. H. 2004, ApJ, 617, 1204 images would help for deriving a surface magnetogram; this Kastner, J. H., Huenemoerder, D. P., Schulz, N. S., Canizares, C. R., & map of the magnetic active regions on the stellar photosphere, Weintraub, D. A. 2002, ApJ, 567, 434 combined with the X-ray light curves, would then allow a self- Kenyon, S. J., Dobrzycka, D., & Hartmann, L. 1994, AJ, 108, 1872 Mason, K. O., Breeveld, A., Much, R., et al. 2001, A&A, 365, L36 consistent three-dimensional model of the corona of AA Tau to Ménard, F., Bouvier, J., Dougados, C., Mel’nikov, S. Y., & Grankin, K. N. 2003, be built. A&A, 409, 163 Taking into account that the duration of the visibility win- Morrison, R., & McCammon, D. 1983, ApJ, 270, 119 dow for AA Tau is about 130 ks per XMM-Newton orbit (2 days), Ness, J.-U., Güdel, M., Schmitt, J. H. M. M., Audard, M., & Telleschi, A. 2004, a full monitoring of two optical eclipse periods of AA Tau A&A, 427, 667 ff Neuhäuser, R., Sterzik, M. F., Schmitt, J. H. M. M., Wichmann, R., & Krautter, (8.22 days) with XMM-Newton is equivalent to an e ective expo- J. 1995, A&A, 297, 391 sure of 1.1 million seconds. This project has the typical duration Ozawa, H., Grosso, N., & Montmerle, T. 2005, A&A, 429, 963 of the large programs anticipated with XMM-Newton for the next Porquet, D., Mewe, R., Dubau, J., Raassen, A. J. J., & Kaastra, J. S. 2001, A&A, decade. 376, 1113 Preibisch, T., Kim, Y.-C., Favata, F., et al. 2005, ApJS, 160, 401 Reale, F., Betta, R., Peres, G., Serio, S., & McTiernan, J. 1997, A&A, 325, 782 Acknowledgements. This research is based on observations obtained with XMM-Newton, an ESA science mission with instruments and contributions di- Rieke, G. H., & Lebofsky, M. J. 1985, ApJ, 288, 618 rectly funded by ESA Member States and NASA. M.F. was supported by the Robrade, J., & Schmitt, J. H. M. M. 2006, A&A, 449, 737 Spanish grants AYA2004-05395 and AYA2004-21521-E. This research was Romanova, M. M., Ustyugova, G. V., Koldoba, A. V., Wick, J. V., & Lovelace, R. V. E. 2003, ApJ, 595, 1009 partly based on data obtained at the 1.5 m telescope at the Sierra Nevada Observatory, which is operated by the Consejo Superior de Investigaciones Scargle, J. D. 1998, ApJ, 504, 405 Científicas through the Instituto de Astrofísica de Andalucía. Schmitt, J. H. M. M., & Favata, F. 1999, Nature, 401, 44 Schmitt, J. H. M. M., & Robrade, J. 2007, A&A, 462, L41 Schmitt, J. H. M. M., Ness, J.-U., & Franco, G. 2003, A&A, 412, 849 Schmitt, J. H. M. M., Robrade, J., Ness, J.-U., Favata, F., & Stelzer, B. 2005, References A&A, 432, L35 Smith, K., Audard, M., Güdel, M., Skinner, S., & Pallavicini, R. 2005, in Anders, E., & Ebihara, M. 1982, Geochim. Cosmochim. Acta, 46, 2363 Proceedings of the 13th Cambridge Workshop on Cool Stars, Stellar Systems Anders, E., & Grevesse, N. 1989, Geochim. Cosmochim. Acta, 53, 197 and the Sun, ed. F. Favata et al., ESA SP-560, 971 Argiroffi, C., Maggio, A., & Peres, G. 2007, A&A, 465, L5 Stassun, K. G., van den Berg, M., & Feigelson, E. 2007, ApJ, 660, 704 Asplund, M., Grevesse, N., & Sauval, A. J. 2005, in Cosmic Abundances as Stelzer, B., & Schmitt, J. H. M. M. 2004, A&A, 418, 687 Records of Stellar Evolution and Nucleosynthesis, ed. T. G. Barnes, III, & Strüder, L., Briel, U., Dennerl, K., et al. 2001, A&A, 365, L18 F. N. Bash, ASP Conf. Ser., 336, 25 Telleschi, A., Güdel, M., Briggs, K. R., Audard, M., & Scelsi, L. 2007, A&A, Audard, M., Skinner, S. L., Smith, K. W., Güdel, M., & Pallavicini, R. 2005, in 468, 443 Proceedings of the 13th Cambridge Workshop on Cool Stars, Stellar Systems Turner, M. J. L., Abbey, A., Arnaud, M., et al. 2001, A&A, 365, L27 and the Sun, ed. F. Favata et al., ESA SP-560, 411 Vuong, M. H., Montmerle, T., Grosso, N., et al. 2003, A&A, 408, 581 Bouvier, J., Chelli, A., Allain, S., et al. 1999, A&A, 349, 619 Walter, F. M., & Kuhi, L. V. 1981, ApJ, 250, 254 Bouvier, J., Grankin, K. N., Alencar, S. H. P., et al. 2003, A&A, 409, 169 Walter, F. M., & Kuhi, L. V. 1984, ApJ, 284, 194 Bouvier, J., Alencar, S. H. P., Boutelier, T., et al. 2007a, A&A, 463, 1017 Wang, J., Townsley, L. K., Feigelson, E. D., et al. 2007, ApJS, 168, 100 Bouvier, J., Alencar, S. H. P., Harries, T. J., Johns-Krull, C. M., & Romanova, White, N. E., Giommi, P., & Angelini, L. 1996, VizieR Online Data Catalog M. M. 2007b, in Protostars and Planets V, ed. B. Reipurth, D. Jewitt, & K. Wilms, J., Allen, A., & McCray, R. 2000, ApJ, 542, 914 Keil, 479 Wolk, S. J., Harnden, F. R., Flaccomio, E., et al. 2005, ApJS, 160, 423 Briggs, K. R., Güdel, M., Telleschi, A., et al. 2007, A&A, 468, 413 Yang, H., Johns-Krull, C. M., & Valenti, J. A. 2005, ApJ, 635, 466 Broos, P. S., Feigelson, E. D., Townsley, L. K., et al. 2007, ApJS, 169, 353 N. Grosso et al.: Enhanced X-ray emission from AA Tauri, Online Material p 1

Online Material N. Grosso et al.: Enhanced X-ray emission from AA Tauri, Online Material p 2

Fig. 2. A subset of X-ray ﬂares from the Chandra Orion Ultradeep Project with peak amplitude and duration larger than the one observed in AA Tau. The left panels show COUP light curves (see Getman et al. 2005b), where blue and red segments indicate the minimum (BBmin)and maximum (BBmax) levels obtained from Bayesian block analysis (Scargle 1998), respectively. The blue dotted lines show the minimum levels. The peak amplitudes are also given. The dashed green lines start 1.7 days after the end of the maximum Bayesian blocks, and indicate the level (2.5 times above the minimum level) that was observed in the third observation of AA Tau. The right panels are an enlargement of the light curve around the Bayesian block showing the peak level and duration. Dots mark the arrival times of individual X-ray photons with their corresponding energies given on the righthand axis. Large vertical gray stripes indicate the ﬁve passages of Chandra through the van Allen belts where ACIS was taken out of the focal plane and thus was not observing Orion. N. Grosso et al.: Enhanced X-ray emission from AA Tauri, Online Material p 3

Fig. 3. EPIC pn (black), MOS1 (red), MOS2 (green) spectra of AA Tau plotted with the same scale. The lines show our best ﬁts using one- temperature plasma combined with photoelectric absorption for observations #1 to #8 (Table 2) from left to right and top to bottom. The residuals are plotted in terms of sigmas with error bars of size one. N. Grosso et al.: Enhanced X-ray emission from AA Tauri, Online Material p 4

Fig. 9. Average OM count rate versus column density. Labels indicate XMM-Newton observations. The crosses mark observations obtained outside the optical eclipse. The dotted hourglass shows the column density of observation #2 when using only one-temperature plasma. N. Grosso et al.: Enhanced X-ray emission from AA Tauri, Online Material p 5

Table A.1. Elemental abundances of the coronal plasma and the absorbing material.

Coronal plasmaa Absorbing materialb

El A(El) n(El)/n(H) angr A(El) n(El)/n(H) angr (1) (2) (3) (4) (5) (6) (7) He 10.99 9.77E-02 1.000 10.93 8.51E-02 0.871 C 8.21 1.63E-04 0.450 8.39 2.45E-04 0.676 N 7.95 8.83E-05 0.788 7.78 6.03E-05 0.538 O 8.56 3.63E-04 0.426 8.66 4.57E-04 0.537 Ne 8.01 1.02E-04 0.832 7.84 6.92E-05 0.562 Na ...... 6.17 1.48E-06 0.691 Mg 7.00 9.99E-06 0.263 7.53 3.39E-05 0.892 Al 6.17 1.47E-06 0.500 6.37 2.34E-06 0.795 Si 7.04 1.10E-05 0.309 7.51 3.24E-05 0.912 S 6.83 6.76E-06 0.417 7.14 1.38E-05 0.852 Cl ...... 5.50 3.16E-07 1.682 Fig. A.1. Comparison of the column density values obtained from spec- Ar 6.30 2.00E-06 0.550 6.18 1.51E-06 0.417 tral fitting when using the old (Anders & Ebihara 1982) and the updated Ca 5.65 4.47E-07 0.195 6.31 2.04E-06 0.892 (Asplund et al. 2005) solar abundances. The dashed line shows the mean Cr ...... 5.64 4.37E-07 0.902 average between the two values of column density. Fe 6.96 9.13E-06 0.195 7.45 2.82E-05 0.602 Co ...... 4.92 8.32E-08 0.967 Ni 5.54 3.47E-07 0.195 6.23 1.70E-06 0.954 sections of Morrison & McCammon (1983) and the (old) solar abundances of Anders & Ebihara (1982). Significant revisions of Notes: Columns (2) and (5) give the element abundances on the loga- rithmic astronomical scale, where the numbers of hydrogen atoms are the solar abundances have been made recently, as a result of the set to A(H) = log n(H) = 12. The numbers of element atoms normal- application of a time-dependent, 3D hydrodynamical model of ized to the number of hydrogen atoms are given in Cols. (3) and (6); the solar atmosphere, instead of 1D hydrostatic models. This has Cols. (4) and (7) compare this ratio to the Anders & Grevesse (1989) decreased the metal abundances, in particular of carbon and oxy- photospheric abundances. gen, which are the main contributors to the photoionization cross a Elemental abundances observed in the coronae of young stars with section above 0.3 and 0.6 keV, respectively. Consequently, the fine X-ray spectroscopy (Güdel et al. 2007a) used in vapec. absolute value of the column density is dependent on the adopted b Elemental solar abundances of Asplund et al. (2005) used in photoelectric absorption model. Using updated solar abundances tbvarabs. is then crucial when the absolute value of the column density is needed (e.g., Vuong et al. 2003). Columns (5)−(7) of Table A.1 Appendix A: Elemental abundances of the coronal give the recent compilation of elemental solar abundances by plasma and the absorbing material Asplund et al. (2005), where the decrease in metal by comparison with Anders & Grevesse (1989) is indicated in the last col- The typical values of plasma element abundances observed in umn. the coronae of young stars with fine X-ray spectroscopy (Güdel We test the impact of these udpated solar abundances on et al. 2007a), which we use for the vapec coronal plasma our fitting by replacing wabs by tbvarabs (Wilms et al. 2000), model in our X-ray spectral fitting with XSPEC,aregivenin which allows us to input new abundances for the absorbing mate- Cols. (2)−(4) of Table A.1. rial. Moreover, tbvarabs also use also updated photoionization The column density of the absorbing material located on the cross sections. We find nearly identical values of temperature line of sight is estimated from spectral fitting using a photoelec- and emission measure; however, as anticipated, the column den- tric absorption model that uses photoionization cross sections, sity value is increased. Figure A.1 shows that column density and solar abundances for the material composition. The wabs values obtained with wabs are underestimated by about 50%. photoelectric absorption model uses the photoionization cross